Desired Performance

The principal scientific goal motivating the
spectrograph design is to obtain in a single exposure 600
spectra of galaxies as faint as the spectroscopic limit of r'~ 18.2
over the three-degree field of the telescope.
As always, the projected
performance is the result of optimizing scientific
return within cost and technology constraints.
We will not explain the compromises (cf. the discussion in
Chapter 1), but will describe
the instruments as we are building them. We
first outline the requirements driving the design;
those that most directly affect it are resolution,
wavelength coverage, number of fibers, and fiber diameter.

Resolution (2000)

The spectral resolution is 2000, roughly where
absorption lines in galaxies and quasars are
resolved. Lower resolution reduces signal-to-noise
by filling in the absorption with continuum light,
higher resolution needlessly reduces spectral range,
given that we can use at most two 2048²
sensors to record the spectrum. Higher resolution
also reduces the signal per pixel so that readout noise
dominates for any but long exposures. The resolution we
have chosen and the efficiencies we expect will give
us shot-noise limited spectra almost everywhere in
the bandpass with exposures as short as 20 minutes.
A resolving power of 2000 is sufficient to distinguish
velocities of 150 km/s in a single data element and
is good enough to measure velocities to better than
&plusmn 20 km/s from an entire galaxy spectrum.

Wavelength Coverage (3900-9100 Å)

Placing the blue limit at 3900 Å ensures
that the H and K lines of CaII are observed
even at zero redshift. To observe H alpha to a
redshift of z = 0.2 or more requires coverage
to at least 8000 Å. Detector size suggests
that going to 9100 Å is possible and appropriate,
and the extra range is valuable for quasars.

Number of Fibers (640)

To finish the survey in five years requires
observing about 550 galaxies in each one
hour exposure with efficiency good enough
to do up to ten exposures on a single long
night. Each field will also have between 50
and 100 quasar candidates plus
a few fibers for the sky and the stars
used for reddening, spectrophotometric
calibration, and correcting for
atmospheric absorption bands in the near
infrared. The design comfortably
accommodates 640 spectra per observation, using
two spectrographs with sensors that are 2048
pixels wide.

Fiber Diameter (180 µm)

For galaxies near our brightness limit, a 3 arcsec
diameter aperture is a good compromise between
collecting most of the galaxy light and rejecting
the night sky. Fibers that preserve f/ratio and
have excellent transmission over our wavelength
range can be obtained in this diameter ( 180 µm).
For galaxies near our faint limit, this
diameter gives a galaxy flux comparable
to the sky.

Technical Constraints

Telescope

The telescope was designed with fiber
spectroscopy in mind so it is hardly
a constraint. The f/5 focal ratio is a good
acceptance beam for optical fibers and the
2.5 m aperture is sufficient to measure
faint galaxy redshifts in reasonable
time. The aperture and f-ratio set the
focal length and hence the physical
diameter of the fibers. The required
input aperture diameter, three arc seconds,
corresponds to 0.18 mm at the focal plane.

CCD Format

The detectors are Tektronix 2048 x 2048
CCDs with square 24 µm pixels. The large
pixels allow a modest demagnification ratio
of 2.5 so the challenge is to put the full
area of the detectors to good use. This
requires a spectrograph camera
with an exceptionally large field of
view (16.5°) and the largest practical beam diameter
(150 mm).

Spectrograph Characteristics

The final design is set after considering
the science requirements, telescope, detector
format, and operational considerations.
The important characteristics are summarized here.
Details appear in following sections.

Measuring the profile of a fiber plug plate in a drilling
template. The template deforms the plate to allow a standard
3-axis milling machine to produce hole axes aligned to the principal
ray axis.

Fiber Optics

Fiber optics is the obvious way to do simultaneous
spectroscopy of hundreds of faint galaxies.
We will use custom drilled aluminum plug-plates
to hold the fibers in the telescope focal plane
(see Figure 7.1).
These plug-plates will be installed in a fixture called
a fiber cartridge and manually stuffed with fibers
before nightfall. The fibers are brought to
the slithead, which is incorporated
in the cartridge and which mates with the
spectrograph. During the night, up to ten
cartridges will be swapped out to observe different fields.

In addition to the object fibers, which are
single strands, there will be ten coherent
fiber bundles capable of imaging a few arcseconds
of the sky. These will be placed on
preselected guide stars and feed a CCD camera mounted
on one of the spectrographs. The guide stars will be used
to center the telescope
on the plug-plate field, adjust the plate scale
of the telescope, control focus, and guide the
exposure. One larger
(30 arcsecond) bundle will be used to measure the sky
brightness. Image quality and photometric data
from these guide and sky bundles will be used to estimate
the exposure time required to complete the observation.

Two Spectrographs

The images of the fiber ends must be adequately
sampled by the detector and be spaced sufficiently
far to prevent crosstalk. Each fiber end
( 180 µm diameter) will just span three pixels
( 3 x 24 µm = 72 µm) with a spectrograph
demagnification of 2.5. Placing the individual fibers
at 360 µm intervals we might
squeeze up to 340 spectra onto a single
2048 x 2048 CCD. In practice, packaging
constraints limit us to 320 spectra per spectrograph
and the fiber ends are spaced at 390 µm intervals.
This is many fewer than the 600 or so needed to
complete the survey in five years so we will
use two identical instruments fed by a
single plug-plate. With 320 fibers in each
instrument, we get 640 spectra per exposure.

Two Channels

If a resolution element (the projected
image of a fiber end) is three pixels
across, then 2048 pixels cover 1700 Å in the
visible at R = 2000 . To cover the
desired spectral range of 3900-9100
Å at the desired resolution of 2000,
we need more than 2048 pixels. Each spectrograph
therefore has two detectors, one covering
3900-6100 Å and the other 5900-9100 Å.
This gives the desired spectral range while
approximating the R = 2000 requirement.
The blue-red split is done with a dichroic
coating on a beamsplitter, the blue side
reflected and the red side transmitted.

Multiple Fiber Cartridges

To maximize efficiency, ten fiber cartridges,
each loaded with a different plug-plate, will
be ready to go at the beginning of the night.
No nighttime plugging is required.

Spectrograph Mounted on Telescope

The spectrographs will be mounted on the
telescope to maintain good
fiber performance during and across exposures.
Having captive fibers routed directly to a
spectrograph on the telescope avoids repeated
bending that would occur if the spectrographs
were on the floor. The disadvantage is that the
spectrographs must not flex as the telescope tracks.

Fiber harness. Twenty fibers are terminated
in the v-groove block at one end have ferrules at the other.

Fiber Feed System

Fibers

We are using silica UV-enhanced step-index fiber with a core diameter of
180 µm and a polyamide protective buffer. The spectrographs are
mounted on the telescope, so each fiber need only be 2 m long. A
sample set of fibers is shown in Figure 7.3. Each
plug-plate has 640 fibers (plus approximately ten fiber bundles for
acquisition and guiding), or 320 per spectrograph. Hence we require
6400 optical fibers plus spares to support ten plug-plate cartridges.

Prototype fiber optic cartridge with 20 fibers. The ends, which are plugged
and unplugged during operations, are protected by tough nylon tubing. The
lens enlarges the v-groove block termination at the slit end.

At the slit, the fibers are separated by 390 µm center to center.
The fiber outside diameter is only 220 µm, so there is insufficient room to
individually terminate the fibers. Our two choices are to mount bare fibers
to the slit or terminate the fibers in sets. Bare fibers are fragile and
are difficult to polish flat and normal to the fiber axis. Mass-terminated
fibers are robust and easily polished flat. We have chosen to terminate the
fibers in v-groove blocks of 20 fibers. The resulting v-groove block is
large enough to be handled easily, yet has few enough fibers that we can
afford to replace a set if a few fibers break. We call a set of 20 fibers a
"harness" (Figure 7.2).

At the plug-plate, each fiber is terminated in a stainless-steel ferrule.
Jacketing on the fiber applies a torque to this ferrule, providing
retention of the ferrule in the plug plate. This approach has been used
successfully in many other plug-plate-based fiber spectrographs. Loss of
light due to tilt of the ferrules in the holes is 0.3% or less. We achieve
this by using high-precision ferrules and holes drilled with high precision
spade drill bits held in a custom-made collet.

The fibers are supported below the plug plate by an anchor block. This
block absorbs stresses induced by plugging and orients the fibers for
maximum retention in the plug plate. Between the ferrule and the anchor
block each fiber is encased in a loose-fitting jacket. The jacket protects
the fiber from undue bending and applies sufficient torque to the ferrule
to retain it in the plug plate.

Between the anchor block and the slit the fibers are not disturbed even
during plugging, and so need not be heavily protected. In most of this
region the fibers will be individually jacketed. Near the slit,
the fibers will be encased in larger
tubing, ten fibers to a tube, to assist in routing the fibers along the
spectrograph slit.

At the spectrograph slit, the twenty fibers of the harness terminate in one
custom-made stainless steel v-groove block (Figure 7.4).
The grooves of
the block are fanned out slightly in such a way that although the block is
polished flat, the beam emitted by each fiber is normal to the slit.

We received and tested prototype harnesses from a number of vendors.
These harnesses have adequate throughput and survive lifetime tests.
The final vendor selection has been made and two complete sets of
640 fibers have been fabricated and tested.

We will install the fiber optic harnesses ourselves. The harnesses are
attached at two points: the anchor block and the slit. The anchor block
mounts to the plug-plate cartridge frame with a single bolt. The v-groove
block is aligned by pressing the output edge against a curved alignment
jig, and attached to the slit with an adhesive. An air-powered fluid
dispenser meters the adhesive to prevent contamination of critical
surfaces. A hole in the slit under each block allows removal of the block
for replacement; a custom-made tool reaches through the hole to push
the block off the slit, while simultaneously holding down the two
neighboring blocks.

Fiber Tester

During manufacturing, each optical fiber will be tested for adequate
throughput using the apparatus shown in Figure 7.5.
White light from an
intensity-stabilized quartz-halogen lamp is fed to the apparatus via a
"source" optical fiber. The end of the source fiber is imaged onto the
fiber under test using a microscope objective, which produces a uniform f/5
converging beam. A microscope eyepiece and pellicle beamsplitter give the
user a view of the fiber under test.

Light from the output end of the test fiber is collected by a pair of
achromatic doublets focussed onto a silicon photodiode. A filter between
the doublets flattens the quartz-halogen spectral curve. A calibrated
aperture blocks light outside a cone of f/4. A computer-controlled
translation stage allows one to accurately locate the appropriate fiber of
the v-groove block in front of the aperture. This same light collection
system may also be used to measure light from the microscope objective,
allowing us to make absolute throughput measurements.

We have constructed two fiber testers. The manufacturer uses one to measure
every fiber and reject those found to have inadequate throughput, and we
use the other to verify the manufacturer's measurements for some fraction
of the fibers.

Optical Design

The spectrograph optical layout is shown in
Figure 7.6. Light from the fiber
optics (A) exits in an f/4 beam, expanded somewhat
from the f/5 input beam of the telescope due
to processes collectively known as focal ratio
degradation. The beam encounters the spherical
collimator mirror (B) and collimated light returns in a
150 mm diameter beam, passes the slit, and
meets the dichroic beamsplitter (C). The blue
light ( < 6000 Å) is reflected while the red
light is transmitted. Past the beamsplitter, the
light encounters the dispersing grism (D).
The dispersed light exits the grism and
meets the wide field f/1.3 240 mm camera that images
the spectra onto the Tektronix 2048 x 2048 CCD.
Optical details of these components are described
here. Mechanical details are in the following section.

Collimator Mirror

The fiber slit is 4.9 inches long, packed with
320 fibers. The obvious collimator design to
use with this long slit is a Schmidt. The Schmidt
is inexpensive to manufacture, has the required
field of view (6°), and uses a mirror, which
makes the instrument compact.

Although we explored a classical Schmidt design,
we found it possible to eliminate the corrector plate
because our imaging requirements did not need the
full Schmidt performance and some of the deficit due
to the missing correctors could be compensated in the
spectrograph cameras. Thus, our final collimator
design is a single mirror with a spherical figure.
The mirror is rectangular, 7 x 17 inches,
with a 49.8 inch radius of curvature. The substrate
is a slumped borosilicate gas fusion blank made by Hextek.

We examined alternatives to the Hextek blank including
aluminum, eggcrate Zerodur, and monoliths. The gas
fusion blank wins easily on weight and cost considerations.
The compromise is in the thermal coefficient of expansion,
which is non-zero but smaller than aluminum. An aluminum
mirror would obviate the need to refocus when the temperature
changes, but is expensive and the coatings are fragile. An
eggcrate Zerodur mirror is immune to large temperature
changes but is expensive and would require the most
refocusing. A monolith would work well enough, but would
not be significantly less expensive than the lighter and
more thermally responsive gas fusion mirror.

An enhanced silver
coating will be placed on the mirror for better than 95%
reflectivity throughout our wavelength range.

Beamsplitter

The beamsplitter, which divides the collimated beam
into red and blue channels, is made of fused silica.
This material has excellent transmission and low dispersion
in our wavelength range. It is thick, 1.5 inches, because
the reflecting surface is a mirror and needs to maintain
a flat figure.

The dichroic coating is efficient in reflection (98%),
reasonably good in transmission (94%), and has a narrow
200 Å crossover range.

Grism

The dispersing elements are grisms with zero angular
deviation at 4960 Å for the blue and 7400 Å
for the red. In our case,
these are right angle prisms with a
transmission grating replicated on the hypotenuse.
While a reflection grating might have been used,
the grism permits mounting the cameras close to
the system pupil, which is about midway on the grating.
With a reflection grating, the cameras have to
be mounted away from the grating to avoid
interference with the incoming beam, making
them larger and more difficult to design.
A plane transmission grating does not work
because the diffracted angle is large,
making geometric losses high (the groove
facets are foreshortened) and forcing the blaze
peak outside the optical band. Our configuration
has little groove shadowing or foreshortening and
results in high grating efficiency.

The ruling densities are 640 and 440 lines/mm for
the blue and red grisms, respectively. Because
master rulings of the size and groove angle needed
did not exist, new masters were ruled.

Camera

Spectrograph camera designed by Harland Epps.
Diameters are in inches. The asphere is on the air surface of the
FPL51Y lens.

The spectrograph camera is an all-transmission
design because the fiber optics generate a
filled beam. The usual practice of hiding
the detector or a secondary mirror in the
center of the beam (as in a Schmidt
camera) cannot be used without unacceptable
light loss.

The cameras are the biggest technical challenge
in the spectrograph. An f/1.3 camera with 240 mm
focal length and 16.5° field of view capable
of using 24 µm pixels cannot be purchased off
the shelf. On the other hand, cameras with
similar performance are now being built primarily
because of the availability of Tektronix 2048
CCDs with 24 µm pixels and the popularity
of fiber spectrographs.

We considered modifying an existing design and
examined the KPNO Bench Spectrograph Camera
(by George Simmons), the VLT spectrograph cameras
(Hans Dekker et al.), and the LRIS and Norris
spectrograph cameras (Harland Epps). All are
Petzval lenses with two widely spaced, approximately
equal powered components. None provided the performance
needed.

We chose to contract Harland Epps (UCSC) to design new
cameras for our spectrographs, shown in Figure
7.7. The result is similar
to the LRIS design but has a smaller diameter and larger
field of view. All but one of the surfaces are spherical
and the asphere is relatively mild.

Detectors

The CCD detectors are thinned Tektronix
2048 x 2048 with square 24 µm
pixels. Readout noise is 5 electrons,
full well is 150,000 electrons. Their
remarkable characteristic is a high
quantum efficiency in the blue spectral
region (see Figure 7.9).

Optical Prescription

Table 7.1 shows the spectrograph optical
prescription.
Some of the surface descriptions can be deciphered from
this example: "Doublet,first,back" refers to the
doublet component, first lens element, back surface
(closest to the CCD). The "Radius"
is the radius of curvature in inches, negative implying
a concave left surface. "Thickness" is the distance
from the current surface to the next;
positive to the right. All materials are from Ohara Glass
except for CaF2 and the lens couplant, Dow-Corning Q2-3067.

Optical Performance

The spectrograph sensitivity is
controlled by the grating throughput
and the CCD quantum efficiency. We are
fortunate to have excellent blue response
in the CCDs. Figure 7.8 shows
the throughput estimate.
This is for photons falling within
the 3 arcsec fiber input aperture and does not include
losses at the telescope mirror coatings or
light that misses the fiber.
Figure 7.9 shows the detail efficiencies
of the spectrograph components.

Although the input apertures cover three pixels
on the detector, we need better
resolution in the optical system to minimize
crosstalk and improve signal-to-noise. The spectrograph
optical design meets this goal nicely, producing rms
spot diameters about the size of a pixel and
resulting in nearly zero crosstalk from adjacent fibers.
Table 7.2 shows rms spot diameters over the field of
the blue CCD. The red performance is similar.
The wavelengths (in A) are listed along the top and the slit
position, measured in mm from the slit center, along the left.
The slits are 4.9 inches high so the
60 mm field position is near the end of the slit, which is
imaged along the edge of the chip.
The wavelength coverage in the blue channel is 3900-6100
A so the table covers one half the area of the CCD.
Performance is excellent over most of the chip.
The r.m.s. spot diameters are about the size of a single pixel.

Scattered light is a notorious problem for straight-through
spectrograph designs and we do not anticipate perfection here.
The gratings will probably be the primary source of scattered
light and unfortunately, there is little we can do to control
this. We will minimize the problem
by careful attention to baffling and cleanliness; note that the
instrument need never be opened to the outside atmosphere
during normal operation. The interior will be purged with dry
nitrogen, preventing condensation and mineral deposition, and
the gas used on all actuators will be exhausted to the exterior.

Mechanical Design

Simulated spectra. Spectra are shown for the red and blue detectors in
A simulated sky-subtracted but otherwise raw spectrum of
a g'=19.8 spiral galaxy at z=0.2 with this system is
shown in Figure 7.10; the top panel shows
the blue half and the bottom the red. The H and K lines of CaII,
the G band, the magnesium feature, and the sodium D lines
are all clearly seen in absorption in this simulated spectrum
of a galaxy considerably fainter than the survey limit.

The telescope optical system is a simple, fast, large field design with a
focal surface flat to 2.6 mm but one where the principal ray deviates from
the normal to the best-focus surface by up to 37 milliradians. For highest
efficiency, the ends of the optical fibers should be positioned on the
best-focus surface with their axes aligned with the principal ray. It
turns out that plug-plate technology can be made to satisfy these criteria
quite nicely.

The plug-plates will be aluminum alloy 2024-T3, 3.2 mm thick and 0.813 m in
diameter. By applying bending moments to the edge of the plate (beyond the
field of view), finite element calculations show that it can be deformed to
match the best-focus surface to an area-weighted 62 microns r.m.s. The
greatest departure from the best-focus surface is 200 microns and occurs at
the center where the images are the best. Overall, the images are not
significantly degraded from the best-focus surface.

As deformed to match the best-focus surface, the hole axes should line up
with the principal ray axes. This is straightforward to accomplish if the
plug-plate is deformed (in the opposite sense) over a properly curved
mandrel, for drilling. If this is done, the drilling can be performed
using a three-axis Computer Numerically Controlled (CNC) milling machine,
i.e., it is not necessary to tilt the drilling head or the plug-plate.

Drill test results indicate that holes can be drilled with an accuracy of
9 microns r.m.s. in position and 4 microns standard deviation in diameter
using short high-precision spade drill bits in a custom-made collet. In
the test, four different bits were used to drill 50 holes each. The
drilling time was 5.8 sec/hole. No significant
degradation in drilling accuracy was
observed for a range of slopes in the work-piece surface from 0 to 70
milliradians.

The plug-plates have a mass of only 4.3 kg. The plug-plates are thin enough
so that the bending stresses, forces and material costs are reasonable.
They are thick enough to provide hole depth adequate to constrain the plug
angular alignment with the hole and to prevent significant gravity-induced
deflections.

Fiber Cartridges

The fiber cartridge consists of a frame that supports the optical fiber
harnesses, spectrograph slithead, and plug-plate holder (see
Figure 7.11).
The plug-plate holder consists of two large rings that warp the plug-plate
to match the telescope best-focus surface. The upper ring
includes a kinematic interface to the telescope to allow repeatable
positioning of the cartridge on the telescope. By assembling the components
into a cartridge, they can be handled as a unit, thereby reducing the
complexity of the plug-plate changing operation. In particular, this
approach addresses the following issues:

The optical fibers are quite fragile and must be protected during
transport to and from the telescope.

The time available to change plug-plates is limited and must be
performed with minimal lighting to avoid affecting neighboring telescopes.

Fiber Cartridges shown with spectrographs.
The large circle is the outline of the instrument rotator.

Ten cartridges are planned, enough so that plugging need not overlap
observations except on the longest, darkest and clearest nights. This,
in turn, allows plugging to be performed by day-shift personnel, is
likely to result in higher reliability and better staff utilization,
allows some work-load leveling, and allows time for the cartridges to
thermalize completely between plugging and observing.

During the day, for each cartridge that was used successfully the previous
night, the plug-plate will be unplugged and removed and an unused
plug-plate will be installed and plugged. On average, over the life of the
survey, about two cartridges per day will need new plug-plates. Our
experiments with a mockup and the experience of other groups with operating
systems indicate that this will take about two person-hours per cartridge
(Limmongkol et al. 1993). Since a full time equivalent (FTE) employee works
only about 4.8 hours per day averaged over all the days in the year, the
labor requirement is about 83% of a FTE, and the operating cost is quite
acceptable. Unfortunately, the work-load is expected to be very uneven,
peaking during dark time in February.

While staffing for this task is likely to evolve, we expect to initially
hire two people with primary responsibility for plugging, with secondary
responsibilities to survey and observatory operations. Other members of the
observatory staff will assist, as needed, in plate changing during peak
load periods.

The correspondence of fibers to the holes in the plug-plate is determined
by a device called the plug plate mapper, which illuminates
each fiber sequentially from the spectrograph ends. The
illuminated fiber will appear as a bright point against a dark background.
A CCD camera will be used to determine the location of each fiber and to
verify its position and throughput. This is a function of
the plugging station, occurs unattended, and should require about 5 minutes.

All cartridge operations occur at the same level, i.e., the telescope
platform and the adjacent support building. The cartridges are assembled
in the plugging room of the support building. They are stored
in a space with doors to both the plugging room and the outside. At night,
the outside door is opened to allow the cartridges to equilibrate to the
temperature of the ambient air. Each cartridge, before use, is moved on a
cart to a holding area near the telescope, where further
equilibration occurs. The telescope is pointed to the zenith for removal or
installation of cartridges. An empty cart is rolled under the cartridge to
be removed. A lifting mechanism built into the fork base lifts the cart
into contact with the cartridge so that the cart is supporting the weight
of the cartridge. The cartridge is detached from the instrument rotator
and the cart/cartridge assembly is lowered to the fork base. The procedure
is reversed to install the new cartridge. Locating surfaces and sockets
guide the cart and cartridge into the proper orientation with respect to
the telescope.

As the cartridge is lifted into place and clamped to the telescope, the
slitheads are simultaneously inserted into sockets in the spectrographs.
The slitheads are attached to the cartridge frame by stiff springs so that
they can move slightly with respect to the rest of the cartridge.
Once the cartridge has been correctly positioned and clamped to the
telescope, the slitheads are loaded against three point kinematic mounts on
the spectrographs by pneumatic clamps. Each slithead will be coded and its
identification relayed to the observer's workstation when it is inserted.
This allows adjustments for each slithead, e.g., image placement on the CCD
and focus, to be made automatically.

Optical Bench

The optical bench maintains the optical alignment.
Figure 7.12 shows how the major optical components
are attached to the optical bench, which is the large
boxy piece in the center.

Optical bench. The large hole
will hold the blue camera. Electronics are seen on the left and
the small box near the bottom center is the slithead pneumatic clamp.

Because the spectrographs are mounted on the telescope,
low flexure is important. We will allow up to 1/10 pixel
flexure at the CCD due to the optical bench during
any one hour exposure.

A finite element engineering model of the
optical bench was used to examine the flexure properties and tune
the design. Our requirement of 1/10 pixel
flexure due to the optical bench in a one hour exposure was
met by a box made of 1/4 inch thick aluminum with
internal reinforcements and an optimized mounting system.
Figure 7.14 shows the flexure performance.
In this test, the instrument was rotated about its long axis and the
position of a spot in an off-axis corner of the CCD was observed.
Figure 7.14 shows the spot position with the spectrograph rotated
60° and turning to -60° at 15° intervals.
The worst case movement through
15° (or a one hour exposure) is considerably less than 1/10 pixel.

Because the telescope focal plane is coupled to the spectrograph
through flexible optical fibers, there is no need to fix the
spectrograph rigidly to the telescope. Instead, mounting fixtures
that accommodate the different coefficients of thermal expansion
of the steel used in the telescope and the aluminum of the spectrograph
will be used.

Slithead Mount

The slitheads are latched into three point kinematic mounts on the
spectrographs and are held in place by pneumatic clamps. Each slithead
will be coded and its identification relayed to the observer's
workstation when it is inserted. This allows slithead-specific
adjustments, image placement on the CCD and focus, to be made.

Hartmann Masks

The spectrograph focus will be determined
using the usual Hartmann test.
Pneumatically actuated collimator
masks take the form of "saloon doors" located
immediately in front of the collimator (these are not shown in
Figure 7.6).

Focus adjustment is required because of
the expansion and contraction of the aluminum
optical bench, which moves the slithead relative
to the collimator mirror. By moving the
collimator to compensate, we keep the slithead
in focus. Temperature sensors will allow us to
read the focus value from a table after initial measurements
are made. Each slithead is expected to have a
slightly different zero-point focus because of manufacturing
variations, so each will be encoded to allow adjustment
for personality.

Collimator Mount

The collimator mount is complex. The collimator must move in and out
(piston) to correct for expansion and contraction of the aluminum
optical bench due to changing temperature. It must also pitch and yaw
with one arcsec precision to allow widening the flat field spectra and
positioning the spectra correctly on the CCD. Small variations between
slitheads require the ability to position accurately the image on the
CCD to avoid, for example, a bad pixel. These motions, which we would
like to control to 1/10 pixel, require linear resolution of one or two
microns.

Our design uses three DC servo motors at the mirror mounting points
driven by an external controller (motors and controller by Physik
Instrumente). A small control computer on the spectrograph commands
the motors through an RS232 link to the motor controller board.

Central Optics

The central optics include the dichroic beamsplitter and the grisms.
Each optic is set in a six-point kinematic mount and the cell, in
turn, is fixed in a three-point mount within the spectrograph optical
bench. Alignment is achieved with tight machining tolerances and
custom machined inserts so no screw adjustments will be necessary.

Camera Cells and Housing

The camera lens cells and housing are adapted directly from Michael
Carr's similar, but larger, LRIS and Norris spectrograph cameras. The
cells follow conventional optical mounting practices except in the
radial supports of the large lenses. Here, six precision machined
inserts of glass-filled Teflon center the lenses on the optical axis.
The more common practice of building a mold of rtv around the lenses
has the disadvantage of not being able to disassemble the camera
easily. Figure 7.15
shows the concept.

Focus and Shutter

The shutter is located immediately in front of the central optics.
Because timing resolution is not critical, we will actuate the
shutters pneumatically, opening or closing them within about 300 ms.

A manual focus mechanism moves the dewars (plus field flatteners) with
respect to the cameras. This adjustment is suitable for initial
camera setup (to correct for red/blue camera differences) and to
accommodate focus changes that might be necessary because of detector
replacements. Note that because the cameras are designed to be
parfocal over our operating temperature range, routine focus of the
spectrographs can be done by adjusting the collimator.

Dewars

The CCDs are mounted in custom dewars using the same ball and socket
adjustment fixture found in the photometric camera. The dewars can be
physically small because they use an autofill liquid nitrogen system
and do not need a large capacity.

Controls and Software

Observer's Program

All spectrograph operations are controlled from a single program
running on the observer's workstation. This program, the observer's
sole interface to the spectrographic system, is a clearinghouse for
observing commands that translates observer's requests into the
commands required by the independent subsystems used for
spectroscopy. These systems include:

Spectrograph microprocessor

Telescope and guider

CCD data acquisition system

Drilling database

Plugging station

Microprocessor

The spectrograph mechanical operations will be controlled by Z-World
Little Giant Z180 microprocessors, one on each spectrograph. The board
was chosen because it is programmable in C and has the right number
and type of I/O ports and A/D converters. They communicate with the
observer's program through an RS-232A serial port. The processor
controls all the mechanical functions on the spectrograph (shutters,
Hartmann masks, collimator tip, tilt, and focus) and monitors ambient
and instrument temperatures for use in adjusting focus and image
positioning.

Telescope

The observer's software communicates fully with the telescope control
system via an Ethernet telnet connection. One might imagine that when
a new plug-plate is locked on the telescope, the spectrograph relays
the plug-plate ID to the observer's software, which looks up the
coordinates in the plug-plate database. The observer's software
commands the telescope to move to the mid HA position for flat field
and wavelength calibration, then to the field for precise positioning,
scaling, and focusing. When these operations are finished, the
observer's software starts the guider, opens the shutters, and begins
observing.

Data Acquisition

The data acquisition hardware for the spectrograph CCDs is a clone of
the photometric camera system. The observer's program actuates the
system through remote procedure calls (RPC) over the Ethernet. The
commands are few and simple, only those required to prep the CCD,
expose, and read out the data. Once the data are off the chips and in
the data system buffer, the images will be downloaded over a high
speed VME link to the observer's workstation where they will be
written to disk. At the end of the night, a tape is written and sent
to Fermilab.

Drilling Database

The drilling database is the list of objects on a plug-plate. It
includes the name and drilling coordinates of each object and the
field coordinates ( alpha , delta , equinox) of the plug-plate. This
information is needed at observing time so the telescope can be
pointed and the plug-plate ID can be incorporated into the data file.
This database is generated at Fermilab from the photometric imaging
data and is delivered on the internet to the observer's workstation at
Apache Point well before the observations.

Plugging Database

To keep track of which fiber went into which plug-plate hole, the
plugging station will be equipped with a device to map the plug-plate
locations to the slithead after plugging is finished. The plugging
database information is written directly to the observer's workstation
disk at plug time. This information is merged with the CCD data
before it is written to disk.

Here is a
possible observing sequence. We anticipate taking both flat field and
wavelength calibrations before every exposure, though the latter may
prove superfluous. These will be obtained at the anticipated mid hour
angle of the exposure to account for the effects of flexure, both in
the spectrograph and the fibers.

For the calibration observations, a white screen is drawn across the
top of the telescope. The screen is illuminated by calibration lamps
mounted on the telescope structure. Flat field exposures will
probably be of two kinds: one is designed to correct for detailed
pixel-to-pixel variations on the detectors. These will be taken with
the collimator mirror nodding slightly (to move the image a few fiber
diameters perpendicular to the dispersion direction) in order to
uniformly illuminate the detector. The others, which might even be
binned vertically (the dispersion direction) during readout, will not
be nodded and are intended to calibrate the fiber throughput as a
function of wavelength. Both flat-field and calibration exposures
will be read in "quick" mode with reduced sample times (but with the
CCD transfer timing the same as for normal readout) to reduce the
readout time to 30 seconds.

After the calibration exposures, the telescope is moved to the
field. Ten coherent fiber bundles feed images of setup stars to a
television camera mounted on the spectrograph. These setup stars are
used to center the field on the plug-plate, adjust the scale of the
telescope (by simultaneously adjusting the primary and secondary
mirrors), and set the telescope focus. Once these operations are
automatically done, the guider is started. A series of four short (1
minute) exposures with the telescope offset 1.5 arcseconds in each
direction in altitude and azimuth build up a 2 x 2 raster of the
plug-plate field. These exposures will be binned heavily in the
dispersion direction on readout; the data will be used to check for
systematic errors in the drilling system and provide
spectrophotometric calibration information. Next, three nominally 15
minute exposures make up the primary spectroscopic observation. Three
exposures allow efficient cosmic ray discrimination, though it is
possible that sufficiently powerful software will work well with two
or even one; we will pursue this issue during the test year.

The sky brightness will be monitored with a guide camera through a
wide-field (30 arcsecond) fiber bundle placed at a blank location in
the field. This information, combined with the transparency and
seeing data available from the ten guide stars will allow us to
control precisely the exposure time needed to achieve the desired
signal to noise ratio. The exposure times will be adjusted for
atmospheric extinction, sky brightness, and galactic extinction with
the aim of producing insofar as possible a uniform limit to the survey
outside the galaxy; the range in exposure times is close to a
factor of two.

Thus the sequence (with estimated timings) is as follows, starting
with the telescope at the zenith between exposures, and assuming that
the previous exposure has just ended:

Remove the old plugplate and install the new one (5 minutes --
this is our current target time, and it appears not unreasonable).

Slew to the position of the anticipated center of the new
exposure. (1 minute) The flat-field screen is being drawn over the
aperture of the telescope during this move.

Take and read the wavelength calibration exposure (1 minute).

Take and read the flat field exposure (1 minute).

Move to the position of the field, acquire the guide stars
and make the necessary adjustments to position, rotation, and scale
(2 minutes).

Take the four offset raster exposures (4 minutes total; the
read time and offsetting time are negligible).

Set back to the center of the field, turn on the guider, and
take and read the first exposure (15 minutes exposure plus 1 minute read).

Second exposure (16 minutes).

Third exposure (16 minutes).

Slew the telescope to the zenith in preparation for the plugplate
change (1 minute).

The total time for this sequence is 63 minutes.

At the end of the series of observations, the data are collected,
merged with the database information that identify the field,
plug-plate, and fiber assignments, and written to disk. At the end of
the night, all of the data are written to a DLT tape and shipped to
Fermilab.